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Accepted for publication in the Astrophysical JournalPreprint typeset using LATEX style emulateapj v. 08/22/09

THE END OF NUCLEOSYNTHESIS: PRODUCTION OF LEAD AND THORIUM IN THE EARLY GALAXY

Ian U. Roederer1, Karl-Ludwig Kratz2, Anna Frebel3, Norbert Christlieb4, Bernd Pfeiffer2, John J. Cowan5,Christopher Sneden1

Accepted for publication in the Astrophysical Journal

ABSTRACT

We examine the Pb and Th abundances in 27 metal-poor stars (−3.1 < [Fe/H] < −1.4) whosevery heavy metal (Z > 56) enrichment was produced only by the rapid (r-) nucleosynthesis process.New abundances are derived from HST/STIS, Keck/HIRES, and VLT/UVES spectra and combinedwith other measurements from the literature to form a more complete picture of nucleosynthesis ofthe heaviest elements produced in the r-process. In all cases, the abundance ratios among the rareearth elements and the 3rd r-process peak elements considered (La, Eu, Er, Hf, and Ir) are constantand equivalent to the scaled solar system r-process abundance distribution. We compare the stellarobservations with r-process calculations within the classical “waiting-point” approximation. In thesecomputations a superposition of 15 weighted neutron-density components in the range 23≤ log nn ≤ 30is fit to the r-process abundance peaks to successfully reproduce both the stable solar system isotopicdistribution and the stable heavy element abundance pattern between Ba and U in low-metallicitystars. Under these astrophysical conditions, which are typical of the “main” r-process, we find verygood agreement between the stellar Pb r-process abundances and those predicted by our model. Forstars with anomalously high Th/Eu ratios (the so-called actinide boost), our observations demonstratethat any nucleosynthetic deviations from the main r-process affect—at most—only the elements be-yond the 3rd r-process peak, namely Pb, Th, and U. Our theoretical calculations also indicate thatpossible r-process abundance “losses” by nuclear fission are negligible for isotopes along the r-processpath between Pb and the long-lived radioactive isotopes of Th and U.

Subject headings: Galaxy: halo — nuclear reactions, nucleosynthesis, abundances — stars: abun-dances — stars: Population II

1. INTRODUCTION

Nucleosynthesis of the heaviest elements in the Uni-verse is accomplished by successive additions of neutronsto existing iron (Fe) group nuclei in stars. Two fac-tors restrict a star’s ability to produce very heavy nu-clei through standard fusion reactions: the endother-mic nature of fusion reactions for species heavier than56Fe and the increased Coulomb barriers that discour-age charged particle reactions in isotopes with sufficientlyhigh Z. Isotopes heavier than those of the Fe group aretherefore overwhelmingly produced by neutron- (n-) cap-ture processes. Either a single neutron can be added ontimescales longer than the mean time to β− decay (theslow [s] n-capture reaction) or many neutrons will beadded before multiple α and β− decays return the isotopeto stability (the rapid [r] n-capture reaction). The rateof neutron captures and the resulting abundance pat-terns are strongly regulated by the physical conditionsand neutron densities at the time of nucleosynthesis.Understanding the production of the heaviest ele-

1 Department of Astronomy, University of Texas atAustin, 1 University Station, C1400, Austin, TX 78712-0259;iur,chris@astro.as.utexas.edu

2 Max-Planck-Institut fur Chemie, Otto-Hahn-Institut, J.-J.-Becherweg 27, D-55128 Mainz, Germany; klk,bpfeiffe@uni-mainz.de

3 Harvard-Smithsonian Center for Astrophysics, 60 Garden St.,MS-20, Cambridge, MA 02138; afrebel@cfa.harvard.edu

4 Zentrum fur Astronomie der Universitat Heidelberg, Lan-dessternwarte, Konigstuhl 12, D-69117 Heidelberg, Germany;N.Christlieb@lsw.uni-heidelberg.de

5 Homer L. Dodge Department of Physics and Astronomy,University of Oklahoma, Norman, OK 73019; cowan@nhn.ou.edu

ments in the Universe provides crucial insight intothe nature of these processes. An exact site forthe location of the r-process has not been conclu-sively identified, although the appearance of r-processmaterial in stars of very low metallicity ([Fe/H] =−3.0)6 argues against possible astrophysical sites (e.g.,neutron-star mergers) that require long evolutionarytimescales (e.g., Argast et al. 2004), and implies in-stead that some association with Type II core collapsesupernovae (SNe) is likely (e.g., Cowan & Thielemann2004; Cowan & Sneden 2006; Farouqi et al. 2009). Thelack of a precise identification of the r-process site,however, has complicated efforts to model it, neces-sitating vast amounts of input nuclear data and in-creasingly more sophisticated model approaches (e.g.,Cowan et al. 1991a; Kratz et al. 1993; Chen et al. 1995;Freiburghaus et al. 1999; Pfeiffer et al. 2001; Kratz et al.2007b; Arnould et al. 2007). Nevertheless, encourag-ing progress has been made in recent years to experi-mentally determine the half-lives, nuclear masses, andn-capture cross-sections for nuclei along the r-processpath (Kratz et al. 1993; Rauscher & Thielemann 2000;Pfeiffer et al. 2001; Kratz et al. 2007b; Schatz 2008), andsome astrophysically-motivated models of the r-processhave been able to reproduce the robust r-processpattern between A ≃ 120 and the actinide region(e.g., Kratz et al. 2007a; Farouqi et al. 2009). Low-and intermediate-mass stars (∼ 1.5–3.0M⊙) that pass

6 We adopt the standard spectroscopic notations that[A/B] ≡ log10(NA/NB)⋆ – log10(NA/NB)⊙ and log ǫ(A) ≡log10(NA/NH) + 12.00 for elements A and B.

2 Roederer et al.

through the asymptotic giant branch phase of evo-lution are the primary source of s-process mate-rial (e.g., Busso et al. 1999; Lattanzio & Lugaro 2005;Straniero et al. 2006), and investigators have had greatsuccess in matching model predictions to observeds-process abundance patterns (e.g., Sneden et al. 2008and references therein). Understanding these processesthen permits the study of additional astrophysics, suchas the use of radioactive isotopes to determine the agesof individual stars and stellar systems (e.g., Cowan et al.1991; Truran et al. 2002).Lead (Pb) and bismuth (Bi) are the two heaviest nu-

cleosynthesis products of the s-process. Pb has fournaturally-occurring isotopes: 204Pb, 206Pb, 207Pb, and208Pb; Bi has one, 209Bi. The 204Pb isotope is blockedfrom r-process production by the stable nucleus 204Hgand will not be considered further here. 208Pb is adouble-magic nucleus (N = 126 and Z = 82, closingboth its neutron and proton shells), which significantlylowers its cross section to further n-capture. One sub-sequent neutron capture (and β− decay) produces thestable atom 209Bi and another produces 210Bi, whichmay either (in its isomeric state) β− decay or (in itsunstable but relatively long-lived ground state, t1/2 =

3×106 yr) α decay. Both processes produce 206Pb. Thus,a Pb-Bi terminating cycle forms (Burbidge et al. 1957;Clayton & Rassbach 1967; Ratzel et al. 2004).In low metallicity stars the ratio of free neutrons to

Fe-peak seeds is higher, and on average more neutronsare captured per seed nucleus (Malaney 1986; Clayton1988). This results in a high fraction of heavy s-processnuclei being produced (Gallino et al. 1998; Busso et al.1999, 2001; Goriely & Siess 2001; Travaglio et al. 2001;Van Eck et al. 2001; Cristallo et al. 2009). Alterna-tively, an extended period of s-process nucleosynthesisin a given star could also produce a large number ofheavy nuclei. These properties of s-process nucleosyn-thesis have led to spectacular Pb enhancements in sev-eral metal-poor stars, such as CS 29526–110 ([Pb/Fe] =+3.3, Aoki et al. 2002), HD 187861 ([Pb/Fe] = +3.3,Van Eck et al. 2003), HE 0024–2523 ([Pb/Fe] +3.3,Lucatello et al. 2003), SDSS 0126+06 ([Pb/Fe] +3.4,Aoki et al. 2008), and CS 29497-030 ([Pb/Fe] = +3.65,Ivans et al. 2005).The termination of the r-process occurs very differ-

ently. Nuclei along the r-process path with A & 250 willundergo (spontaneous, neutron-induced and β-delayed)fission and repopulate the r-process chain at lower nu-clear masses (Panov et al. 2008). Neutron-rich nu-clei produced in the r-process that have closed neutronshells at N = 126 (at the time of the termination ofthe r-process) will rapidly β− decay to atoms of ele-ments at the 3rd r-process peak—osmium (Os), iridium(Ir), and platinum (Pt), but not Pb. The nuclei thatform with A = 206–208 nucleons in the r-process (i.e.,those that will β− decay to the Pb isotopes) have nei-ther a magic number of neutrons or protons, so theyare produced in smaller relative amounts. Even in themetal-poor stars with the most extreme r-process over-abundances, the [Pb/Fe] ratios are extremely depressed.Pb produced exclusively in the r-process has been de-tected previously in only three stars with [Fe/H] <−2.0 (HD 214925, [Fe/H] = −2.0, Aoki & Honda

2008; HD 221170, [Fe/H] = −2.2, Ivans et al. 2006;CS 31082–001, [Fe/H] = −2.9, Plez et al. 2004). Thisoverall lack of observational data concerning Pb abun-dances in very metal-poor, r-process enriched stars servesas one motive for our present study.There is an additional nucleosynthesis path for three

of the Pb isotopes. All nuclei heavier than 209Bi areultimately unstable to α or β decay and will follow de-cay chains that leave them as either Pb or Bi. In thisway, the Pb and Bi abundances in a star enriched byr-process material will increase with time as these heav-ier nuclei gradually decay. Only two isotopes of thorium(Th) and uranium (U), 232Th and 238U, have halfliveslonger than 1 Gyr (t1/2[

232Th] = 14.05± 0.06 Gyr and

t1/2[238U] = 4.468± 0.003 Gyr; Audi et al. 2003 and ref-

erences therein). In old, metal-poor stars, these are theonly heavier atoms that we have any hope of observingtoday. Only one detection of Bi has ever been madein a metal-poor star, in the strongly s-process-enrichedstar CS 29497-030 (Ivans et al. 2005). This transition ofBi i (3067 A) lies very near the atmospheric transmissioncutoff in a very crowded spectral region. Similarly, Uis very difficult to detect in metal-poor stars owing toits relatively small abundance and the fact that the oneuseful transition of U ii in the visible spectrum (3859 A)is severely blended with CN molecular features and lieson the wing of a very strong Fe i line. Measurementsof the U abundance have only been reported for threemetal-poor stars.Rather than focus on Bi or U, we investigate Pb and Th

to better understand the heaviest long-lived products ofr-process nucleosynthesis. In this study we examine thePb and Th abundances in a sample of 47 stars, derivingnew abundances for 14 stars and adopting literature val-ues for the others. We compare our results with r-processnucleosynthesis model predictions and discuss their im-plications for stellar ages and the chemical enrichment ofthe Galaxy.

2. OBSERVATIONAL DATA

For 12 stars in our sample, we employ the same spectraused by Cowan et al. (2005). These data were collectedwith the Space Telescope Imaging Spectrograph (STIS)on the Hubble Space Telescope (HST ) and the High Res-olution Echelle Spectrograph (HIRES; Vogt et al. 1994)on the Keck I Telescope. Most, but not all, of these12 stars are covered in both sets of observations; Ta-ble 1 indicates which stars were observed with thesetwo facilities. The STIS spectra cover the wavelengthrange 2410 < λ < 3070 A at a resolving power R ≡λ/∆λ ∼ 30,000 and S/N & 50/1. The HIRES spectracover the wavelength range 3160 < λ < 4600 A at a re-solving power R ∼ 45,000 and 30/1 . S/N . 200/1.We also derive abundances from two stars ob-

served with UVES on the VLT, analyzed previouslyby Hayek et al. (2009). These spectra cover 3300 <λ < 4500 A at resolving powers of R ∼ 57,000(CS 29491–069) and R ∼ 71,000 (HE 1219–0312).We supplement our sample with measurements from

the literature. We include 5 stars from Johnson & Bolte(2001) and Johnson (2002), 11 stars from Aoki & Honda(2008), and one star each from Hill et al. (2002) andPlez et al. (2004), Honda et al. (2004), Christlieb et al.

Lead and Thorium from the r-process in Metal-Poor Stars 3

(2004), Ivans et al. (2006), Frebel et al. (2007), andLai et al. (2008). We also include recent Pb or Th abun-dances for the globular clusters M5 (2 stars; Yong et al.2008a,b), M13 (4 stars; Yong et al. 2006), M15 (3 stars;Sneden et al. 2000), and M92 (1 star; Johnson & Bolte2001 and Johnson 2002). Th has also been measured inone star in one dwarf spheroidal (dSph) galaxy, Ursa Mi-nor (UMi; Aoki et al. 2007). Including these additionalstars increases our sample to 47 stars.

3. ABUNDANCE ANALYSIS

For the 14 stars whose abundances are reexam-ined here, we adopt the model atmospheres de-rived by Cowan et al. (2002), Sneden et al. (2003),Simmerer et al. (2004), Cowan et al. (2005), andHayek et al. (2009). These parameters are listed in Ta-ble 1. We perform our abundance analysis using themost recent version of the 1D LTE spectral analysiscode MOOG (Sneden 1973). We compute model atmo-spheres without convective overshooting from the Kurucz(1993) grid, using interpolation software developed byA. McWilliam and I. Ivans (2003, private communica-tion).In addition to Pb and Th, we also derive abun-

dances of lanthanum (La), europium (Eu), erbium (Er),hafnium (Hf), and Ir. Our goal is to sample then-capture abundance ratios at regular intervals in N(or Z) using species that can be relatively easily andreliably derived from ground-based blue and near-UVspectra. Abundances derived from individual transi-tions are listed in Tables 2 and 3, along with the rele-vant atomic data. Some studies predate the most recentlaboratory measurements of n-capture transition proba-bilities, so we update the abundances from those stud-ies. We adopt the latest log(gf) values for the n-capturespecies from a number of recent studies: La (Lawler et al.2001a), Eu (Lawler et al. 2001b), Er (Lawler et al. 2008),Hf (Lawler et al. 2007), Ir (Ivarsson et al. 2003), Pb(Biemont et al. 2000), and Th (Nilsson et al. 2002). Fi-nal abundances for all elements, including updates of theliterature values, are listed in Table 4.

3.1. Lead

We examine three Pb i lines in each of our 14 stars:2833, 3683, and 4057 A. The 2833 A resonance line is thestrongest transition in stellar spectra, which demands useof the STIS UV spectra; all other transitions are acces-sible from spectra obtained with ground-based facilities.Only one of the naturally-occurring Pb isotopes, 207Pb,

exhibits hyperfine structure due to its non-zero nuclearspin. Simons et al. (1989) measured the isotope energyshifts and relative hyperfine splittings from FTS spec-tra. We adopt their measurements and the Solar System(S.S.) isotopic fractions (Lodders 2003) in our syntheses.The s-process produced the majority of the Pb in the S.S.In our stellar sample, where the r-process is the domi-nant contributor to the n-capture material, one mightexpect that a very different isotopic mix could affect ourderived Pb abundances. We find, however, that no sen-sible variations in the isotopic fractions can be detectedin any of our Pb lines due to the relatively small isotopeshifts (. 0.02–0.03 A) and overall weakness of the lines.The 2833.053 A Pb i line is blended with an Fe ii line at

2833.086A. A log(gf) value for this line is listed in the

critical compilation of Fuhr & Wiese (2006), who quotea “C” accuracy in log(gf); i.e., ±25% or about 0.1 dex.Another unidentified line at 2832.91 A mildly blends withthe blue wing of the Pb line, but this has little effecton the derived Pb abundance. Both of these blendingfeatures are weak in stars with [Fe/H] < −2.0 and arenegligible in the most metal-poor stars of our sample.The 3683.464 A Pb i line sits on the wing of a Balmerseries transition. We empirically adjust the strength ofthis Balmer line to reproduce the continuous opacity atthe location of the Pb line. An Fe i line at 3683.61A onlyslightly blends the red wing of the Pb line; we derive alog(gf) value of −1.95 for this line from an inverted solaranalysis. The 4057.807A Pb i line is heavily blended withCH features. In typical r-process enriched stars (with so-lar or subsolar C/Fe and approximately solar Pb/Fe), thePb may be marginally detectable; if a star has supersolarC it is nearly impossible to identify absorption from Pb.We measure the Pb abundance in five stars in our sam-

ple, all from either the 2833 or 3683 A lines. For allnon-detections we report upper limits on the Pb abun-dance. These individual measurements are listed in Ta-bles 2 and 3. In Figure 1 we display these three Pb linesin HD 122956, where we derive an abundance from the2833 and 3683 A lines and an upper limit from the 4057 Aline. Our choice of continuum normalization most affectsthe abundance of the 2833 A line due to the relative lackof line-free continuum regions nearby. Our stated uncer-tainties account for uncertainties in the continuum place-ment.

3.2. Thorium

We examine four Th ii lines in each of our 14 stars:3539, 4019, 4086, and 4094 A. All of these transitionsarise from the ground state. The 4019.129A Th ii line isrelatively strong but suffers from a number of blends. Ex-tensive reviews of the blending features have been madepreviously (e.g., Lawler et al. 1990; Morell et al. 1992;Sneden et al. 1996; Norris et al. 1997; Johnson & Bolte2001). We provide only minor updates to these analyses.An Fe i line sits at 4019.04 A; Fuhr & Wiese (2006) re-port a log(gf) for this line, scaled fromMay et al. (1974).They report an “E” accuracy for this line; i.e., ± > 50%or & 0.2–0.3 dex. We allow the strength of this line tovary within these amounts to match the observed spec-trum just blueward of the Th line. Lawler et al. (1990)determined the log(gf) for a Co i line at 4019.126 A,though this line was not found to contribute significantlyin the very metal-poor stars examined here. The strengthof the hyperfine split Co i line at 4019.3 A can be treatedas a free parameter to match the observed spectrum justredward of the Th line. Sneden et al. (1996) identified aCe ii blend at 4019.06 A in CS 22892–052 that may ex-plain extra absorption in the blue wing of the Th line.We adopt an empirical log(gf) = −0.10 for this line,which matches the observed spectrum in BD+17 3248.The Nd ii line at 4018.82 A, for which Den Hartog et al.(2003) report an experimental log(gf) value, can be usedto estimate the Nd abundance. We then set the Ce abun-dance from the typical Nd/Ce ratio in r-process enrichedmetal poor stars, [Nd/Ce] ≈ +0.25. In stars with anextreme overabundance of n-capture material, such as

4 Roederer et al.

Fig. 1.— Syntheses of the three Pb i lines in HD 122956. Theobserved spectrum is indicated by black squares. In the top twopanels, our best-fit synthesis is indicated by the blue line. Changesto this synthesis of ± 0.30 dex are indicated in red. A synthesiswith no Pb present is indicated by the black line. We only derivean upper limit for Pb from the 4057 A line, indicated by the greenline in the bottom panel.

CS 22892–052, Sneden et al. (1996) note that exclusionof the Ce blend would only increase the Th abundanceby ≈ 0.05 dex. This would almost certainly decrease instars with less severe n-capture overabundances.The final serious contaminant to the 4019 A Th line is

the B2Σ− −X2Π(0–0) 13CH transition doublet (P11[16]and P22[16]; Kepa et al. 1996), which was first discussedby Norris et al. (1997). It is thus necessary to estimatethe C abundance and the 12C/13C ratio in each of ourstars. The corresponding 12CH doublet lies approxi-mately 1 A redward of the 13CH doublet. Norris et al.(1997) have noted that the wavelengths of these featuresmay be incorrect in the Kurucz lists by ∼0.15–0.25A.We determine the absolute position of each of these linesfrom the 12CH absorption lines, but we do not changethe isotope shift between the 12CH and 13CH features.We empirically set the overall strength from the 12CHdoublet, adjusting the log(gf) values for the 12CH tran-sitions and the 13CH transitions together. We then em-

Fig. 2.— Syntheses of the four Th ii lines in BD+17 3248. Sym-bols are the same as in Figure 1. We derive a Th abundance fromeach of these four lines in this star.

ploy these same steps with the 13CH and 12CH dou-blets at ≈ 4006 and 4007 A, respectively (P11[15] andP22[15]). The

13CH doublet is relatively unblended here,permitting measurement of the 12C/13C ratio. The over-all strength of the contaminating 4019 A 13CH featurecan then be reduced by the 12C/13C ratio. Results de-rived from this method agree well with results derivedfrom the CH linelist of B. Plez (2007, private commu-nication). Our measurements of 12C/13C are listed inTable 5.The 3539.587A Th ii line is relatively weak but un-

blended. The red wing of the 4086.521 A Th ii linemarginally blends with the blue wing of a strong La iifeature at 4086.71 A. The blue wing of the 4094.747 ATh ii line blends with several features, including a Nd iiline at 4094.63 A, an Er ii line at 4094.67 A, and a 12CHline at 4094.70 A. Reliable log(gf) values are not knownfor any of these features. The strength of these blendscan be adjusted empirically in most stars; in other caseswe instead measure an upper limit for the Th abundance.A Ca i line at 4094.93 A also mildly blends the red wingof the Th line, but this line can be easily accounted for inour syntheses. Syntheses for these four lines are shown

Lead and Thorium from the r-process in Metal-Poor Stars 5

in Figure 2.

3.3. Uncertainties and Comparisons to Previous Studies

Pb is one of several n-capture species observed in theoptical regime in the neutral state in metal-poor stellaratmospheres, due to its relatively high first ionization po-tential (I.P.), 7.42 eV. The first I.P. of atoms of the 3rd

r-process peak (e.g., Os, Ir, Pt) is even higher (8.35, 8.97,and 8.96 eV, respectively), and these species, too, are ob-served in the neutral state in metal-poor stars. Th, witha lower I.P. of 6.31 eV, is observed in its singly-ionizedstate, as are all of the rare earth species. The singly-ionized states of these atoms are the dominant species intypical metal-poor stellar atmospheres—even for Pb andthe 3rd r-process peak elements—but less so than Th orthe rare earth elements. Thus the abundances derivedfor Th and the rare earth elements are determined fromthe majority species, but we caution that this is not thecase for Pb, where most of the atoms are not in the neu-tral state. Further exploration of this issue is beyond thescope of the present work, but we would welcome moredetailed atomic model calculations of the Pb ionizationbalance.Species in different ionization states clearly respond

differently to conditions in the stellar atmosphere.Cowan et al. (2005) present an extended discussion ofthe uncertainties between ratios of elements with differ-ing ionization states. To summarize those results, foratmospheric uncertainties of ∆Teff = ±150 K, ∆ log g =±0.3 dex, and ∆vt = ±0.2 km s−1, the total uncertain-ties in ratios between neutral and singly-ionized species(e.g., Pb/Eu or Th/Pb) are typically ≃ ±0.20 dex. Un-certainties in ratios between species of the same ion-ization state (e.g., Th/Eu) are much smaller, typically. ±0.05 dex. When considering ratios between species ofdifferent ionization states, we add an additional 0.20 dexuncertainty in quadrature with the individual measure-ment uncertainties.Uncertainties in the abundance ratios resulting from

neglect of the electron scattering contribution to the con-tinuous opacity should be small since these ratios are de-rived from transitions in the same spectral region. Whilethe continuous opacity may be affected more strongly byelectron scattering at the Pb i transition in the ultra-violet, the abundances derived from this line generallyagree with those derived from the redder lines in individ-ual stars in our sample.We compare our derived Pb, Th, and 12C/13C for four

stars with previous high-resolution analyses in Table 6.Pb abundances have been derived previously for onlytwo of the five stars in which we derived a Pb abun-dance, and in both cases we agree within the uncertain-ties. Our Th abundances in BD+17 3248, CS 22892–052,CS 29491–069, and HE 1219–0312 are in good agreementwith previous studies. The previously reported Th abun-dances in HD 6268, HD 115444, and HD 186478 exhibita great degree of scatter, but our abundances are consis-tent with these measurements.Figure 3 compares our 12C/13C ratios with those de-

rived in Gratton et al. (2000) for metal-poor stars on themain sequence, lower red giant branch (RGB), and upperRGB. The individual luminosities for stars in our sam-ple are also shown in Table 5. The majority of stars inour sample lie on the upper RGB and all have 12C/13C

Fig. 3.— 12C/13C ratios as a function of luminosity. Filledblack squares represent measurements from Gratton et al. (2000)and open black triangles represent lower limits for stars with[Fe/H] < −1.0. Filled red squares and open red triangles indi-cate our measurements and lower limits, respectively. Evolutionaryclassifications from Gratton et al. (2000) are indicated.

in good agreement with the Gratton et al. (2000) starsat similar luminosities. The dredge-up processes thatmoderate the decreasing 12C/13C ratio with increasingluminosity have no effect on the n-capture abundancesin these stars.To explore the systematic uncertainties present when

mixing abundances from different studies, we have com-pared the log ǫ (Eu) abundance and log ǫ (La/Eu) ra-tio derived in the present study to Honda et al. (2004)(5 stars in common), Sneden et al. (2009) (3 stars),Aoki & Honda (2008) (2 stars), and Hayek et al. (2009)(2 stars). Sneden et al. (2009) used the same spec-tra, model atmosphere grid and parameters, and anal-ysis code to derive abundances as we have, differingonly in the list of lines and the “human element”present when different investigators make the same mea-surement. Hayek et al. (2009) used the same spectraand model atmosphere parameters as we have, differ-ing in all other components. Honda et al. (2004) andAoki & Honda (2008) have used the same grid of modelatmospheres as we have, but we have no other compo-nents of our analysis in common. All five of these stud-ies employed spectral synthesis techniques to derive theabundances of La and Eu.We find negligible offsets with respect to Sneden et al.

(2009), ∆log ǫ (Eu) = −0.03±0.03 and ∆log ǫ (La/Eu) =−0.01 ± 0.02. We find moderate and significant off-sets with respect to Honda et al. (2004), Aoki & Honda(2008), and Hayek et al. (2009) in ∆log ǫ (Eu), −0.09±0.07, +0.14 ± 0.13, and +0.08 ± 0.01 dex, respectively.This is not unexpected given all of the different com-ponents that enter into the derivation of an elementalabundance. What is surprising, perhaps, is that sig-nificant differences are also found when comparing thelog ǫ (La/Eu) ratio, which should be largely insensitive todifferences in the model atmosphere grid and parameters.The differences are largest with respect to Honda et al.(2004) and Aoki & Honda (2008), −0.11 ± 0.02 and−0.10 ± 0.01, respectively; a smaller difference is foundwhen comparing with Hayek et al. (2009), +0.04± 0.01.This suggests, perhaps, that a significant source of thedifference arises from the lines used and the algorithmfor reduction and continuum normalization of the stel-lar spectrum, although one should be somewhat cau-

6 Roederer et al.

tious about over-interpreting these differences with onlytwo stars in common between our study and each ofAoki & Honda (2008) and Hayek et al. (2009).

4. THE R-PROCESS NATURE OF OUR SAMPLE

In order to correctly interpret the Th/Pb ratios inthese stars, it is important to demonstrate that ther-process has been the only source of their n-capturematerial. Even small contributions from the s-processwill very easily bias the derived Pb abundances inr-process enriched stars. At low metallicity, an increasein the Pb abundance is one of the earliest signatures ofs-process nucleosynthesis. Th can only be produced inthe r-process and is unaffected by s-process contribu-tions.The equivalence between the relative distributions

of abundances for 56 ≤ Z ≤ 79 and the predictedr-process contribution to these species’ abundances inS.S. material, seen in a growing number of metal-poorstars, is clear evidence that the r-process has beenthe only significant source of n-capture material inthese stars (Sneden et al. 2008, 2009). These stars in-clude BD+17 3248 (Cowan et al. 2002), CS 22892–052(Sneden et al. 1996, 2003), CS 31082–001 (Hill et al.2002), HD 115444 (Westin et al. 2000), HD 221170(Ivans et al. 2006), and HE 1523–0901 (Frebel et al.2007). Sneden et al. (2009) have recently remeasuredand/or updated the rare earth (i.e., 57 ≤ Z ≤ 72)abundances in five of these stars. Since the Th andU abundances of CS 31082–001 are enhanced relative tothe rare earths, for now we will exclude this star fromthe set of standards. A complete chemical analysis ofHE 1523–0901 is underway (A. Frebel et al., in prepara-tion). We accept the remaining four stars as the templatefor “standard” r-process enrichment.In the r-process, the La/Eu ratio is ≈ 1.5, whereas in

the s-process the La/Eu ratio is ≈ 56. In S.S. material,about 69% of the La originated in the s-process, whereasonly about 5% of the Eu originated in the s-process(Sneden et al. 2008, with updates from Gallino). Fur-thermore, La and Eu are two elements with multiple ab-sorption features in the spectra of metal-poor stars, andthe log(gf) values for these transitions are well-known,so their abundances can be derived with minimal line-to-line scatter. The La/Eu ratio is an excellent discrimi-nant of the relative amounts of s- and r-process materialpresent in these stars. In Figure 4 we show the La/Euratios as a function of [Fe/H] for our entire sample. Forcomparison, the pure s-process and pure r-process nucle-osynthesis predictions for the La/Eu ratio and the S.S.ratio are also shown.It is clear that the measured La/Eu ratios in these

stars lie close to the pure r-process predictions—but howclose? The mean log ǫ (La/Eu) ratio for the four stan-dard stars is +0.18 ± 0.03 (σ = 0.06), ranging from+0.09 (CS 22892–052) to +0.23 (BD+17 3248). We con-servatively estimate that any star with log ǫ (La/Eu) ≥+0.25 has a non-negligible amount of s-process mate-rial present. Assuming log ǫ (La/Eu)pure−r = +0.18 andlog ǫ (La/Eu)pure−s = +1.75, this limit identifies starswith no more than ≈ 0.5–1.0% of their n-capture mate-rial originating in the s-process. This explicitly assumesthat the four r-process standard stars contain no amountof s-process material. The range and uncertainties of

the La/Eu ratios in these stars set the limit of our abil-ity to determine this percentage. Even if the actual purer-process La/Eu ratio was 0.1 dex lower than our mean—roughly equivalent to the La/Eu ratio in CS 22892–052,which has the lowest La/Eu ratio of any of our standardstars—this limit would still represent only a 1.1% contri-bution from the s-process. According to this definition,27 stars in our sample have been enriched by only ther-process, which we refer to as the “r-process-only” sam-ple for the remainder of this paper.Elemental abundances for n-capture elements are usu-

ally sums over multiple naturally-occurring isotopes ofthese species. Sneden et al. (2002), Aoki et al. (2003),Lundqvist et al. (2007), and Roederer et al. (2008) haveshown that several of the stars in our r-process-only sam-ple have samarium and Eu isotopic mixes consistent withr-process nucleosynthesis. This lends further credibilityto our assertion that the n-capture material in stars inour sample originated only in the r-process.Now consider the stars in our sample with just a small

amount of s-process material, those with log ǫ (La/Eu)≥+0.25. The La/Eu, Pb/Eu, and Pb/La abundances forthese stars are shown in Figure 5. While the La/Euratio shows no evolution with [Fe/H], the Pb/La ratiodisplays a marked increase with decreasing metallicity.This demonstrates that when the s-process operates atlow metallicity a relatively large amount of material ac-cumulates at the 3rd s-process peak due to the higherneutron-to-seed ratio. Furthermore, this effect is notice-able in stars where only ≈ 2.0% of the n-capture ma-terial originated in the s-process (as determined fromthe La/Eu ratios of the stars with Pb/La ratios dis-played in the middle panel of Figure 5). In contrast,in stars with log ǫ (La/Eu) < +0.25 the Pb/La ratiodisplays no trend with metallicity. The combination ofthese two facts reinforces our assertion that by choosinglog ǫ (La/Eu)r < +0.25 we have identified a sample ofstars that are free of any detectable traces of s-processenrichment.In Figures 4 and 6 we examine the Er/Eu, Hf/Eu,

and Ir/Eu ratios. By comparing these abundance ra-tios to those found in the four r-process standard stars—whose compositions have been analyzed in excruciatingdetail for all of the n-capture species with accessibletransitions—we can further characterize and establishthe r-process-only nature of our sample. Stars with ahint of s-process material (identified by their La/Eu ra-tios) are marked in blue, and stars with only r-processmaterial are marked in red. The r/Eu ratios in stars withan r-process-only signature are consistent with a singlevalue. The mean stellar Hf/Eu ratio differs somewhatfrom the predicted r-process-only ratio, suggesting thatthe predicted S.S. breakdown for Hf may need minor revi-sions (Lawler et al. 2007, Sneden et al. 2009; see furtherdiscussion in § 6). The mean ratios for the r-process-onlystars are listed in Table 7. The r/Eu ratios are constantfor La (A = 139), Er (A = 162–170), Hf (A = 174–180),and Ir (A = 191–193), extending through the entirety ofthe rare earths and to the 3rd r-process peak.Figure 7 displays the Pb/Eu and Th/Eu ratios. The

Th/Eu ratio for nearly all stars is consistent with asingle value over the entire metallicity range of themeasurements. This remarkable correlation is not just

Lead and Thorium from the r-process in Metal-Poor Stars 7

Fig. 4.— Comparison of the La/Eu and Er/Eu ratios in our sample. Red squares represent measurements in stars with log ǫ (La/Eu) <+0.25, and blue squares represent measurements in stars with log ǫ (La/Eu) ≥ +0.25. Black dots in the middle of these points signifymeasurements made in the present study. In the upper panels, the dotted black line represents a 1:1 ratio of the elemental abundancesand the red dashed line represents the mean ratio of the four “standard” r-only stars in Table 7. In the lower panels, the red dottedline represents the pure-r-process nucleosynthesis prediction (Sneden et al. 2008, with updates from Gallino), the blue short dashed linerepresents the pure-s-process nucleosynthesis prediction, and the green long dashed line represents the S.S. meteoritic ratio (Lodders 2003).

seen in the field stars, but also in globular clustersand one star in a dSph system. Despite our effortsto detect Pb in low-metallicity stars enriched by ther-process, only one convincing detection exists below[Fe/H] = −2.2, CS 31082–001 ([Fe/H] = −2.9, Plez et al.2004). HE 1523–0901 ([Fe/H] = −2.95, Frebel et al.2007) has a Th/Eu ratio consistent with the standardr-process-only stars, yet the Pb upper limit derived byFrebel et al. (2007) indicates that the Pb in this starlies at least 0.3 dex below the r-process value seen inthe r-process-enriched stars with −2.2 < [Fe/H] < −1.4.CS 22892–052 ([Fe/H] = −3.1, Sneden et al. 2003) has aPb upper limit that is nearly identical to the Pb abun-dance found in the stars with −2.2 < [Fe/H] −1.4. Thesethree stars have approximately the same metallicity andhigh levels of r-process enrichment ([Eu/Fe] = +1.6,+1.8, and +1.6, respectively). For the majority of metal-poor stars with known Pb and Th abundances, ther-process pattern appears to continue to the actinidesas well. A few notable exceptions are discussed in Sec-tion 5.

5. NUCLEOSYNTHESIS OF PB AND TH IN THER-PROCESS

Four stars in the r-process-only sample have Th/Euratios elevated by ∼ 0.3–0.4 dex relative to the otherstars in the r-process-only sample: CS 30306–132,CS 31078–018, CS 31082–001, and HE 1219–0312. This“actinide boost” (Schatz et al. 2002) describes the en-hanced Th (and, in CS 31082–001, U) abundance ra-tio(s) relative to the rare earth elements. All four starshave metallicities in the range −3.0 < [Fe/H] < −2.5and have high levels of r-process enrichment ([Eu/Fe] ≥+0.85, including three with [Eu/Fe] > +1.2), but sev-eral other stars with no actinide boost also have similarlevels of r-process enrichment (e.g., CS 22892–052 andHE 1523–0901). The distinction between these four starsand the remainder of the r-process sample is rather clean,with no stars having −0.40 < log ǫ (Th/Eu) < −0.26.In Figure 8 we compare the mean abundances of these

four stars to those of the other r-process-only stars (alsosee Table 7). The S.S. r-process pattern is shown forreference. The abundances are normalized to the S.S.Eu abundance. The mean La, Er, and Ir abundancesbetween the two groups are identical; the mean Hf abun-

8 Roederer et al.

Fig. 5.— La/Eu and Pb/La ratios in stars with a hint ofs-process enrichment and stars with only r-process enrichment.Blue squares represent stars with log ǫ (La/Eu) ≥ +0.25, while redsquares represent stars with log ǫ (La/Eu) < +0.25. The lines rep-resent least-squares fits to the points. The top panel demonstratesthat neither of these two groups of stars exhibits any significantslope in La/Eu with [Fe/H]. It is clear that the Pb/La ratio in-creases with decreasing metallicity in stars with a small s-processcontribution (middle panel), while no such increase is discerniblein stars with log ǫ (La/Eu) < +0.25 (bottom panel). The starCS 31082–001 ([Fe/H] = −2.90, log ǫ (Pb/La) = +0.07) has beenexcluded from the fit in the bottom panel; see Section 5 for details.

dance in stars with the actinide boost is determined fromonly one star, so we do not regard the 0.10 dex discrep-ancy as significant. The mean Pb abundance in starswith the actinide boost is also derived from only one star,but it differs from the mean of the standard r-process-only stars by 0.51 dex. The abundances of the stars withthe actinide boost are identical to the abundances of thestars without an actinide boost from the rare earth do-main to the 3rd r-process peak; any differences in the nu-cleosynthetic process(es) that produced their n-capturematerial affect—at most—only the region beyond the 3rd

r-process peak.In Figure 9 we show Pb/Eu ratios for the r-process-

only sample. These stellar ratios are compared withour predictions, made using the classical waiting-pointassumption—defined as an equilibrium condition be-tween neutron captures and photodisintegrations—forthe r-process, employing the Extended Thomas-Ferminuclear mass model with quenched shell effects farfrom stability (i.e., ETFSI-Q, Pearson et al. 1996). Al-though this approach makes the simplifying assump-tions of constant neutron number density and temper-ature as well as instantaneous nuclear freezout, the equi-librium model calculations reproduce the S.S. abun-dances well; see, e.g., Kratz et al. (1993), Cowan et al.(1999), Freiburghaus et al. (1999), Pfeiffer et al. (2001),and Kratz et al. (2007a). These calculations are model-independent and look only at the astrophysical condi-tions that lead to the r-process and reproduce the S.S.r-process abundances. These waiting point calculationshave been confirmed by more detailed dynamic (i.e., non-equilibrium) network calculations (Farouqi et al. 2009).Our approach can be considered reliable only if weachieve a “consistent” picture—meaning that the abun-dances are solar—with logical astrophysical assumptionsfor the three heaviest r-process “observables;” i.e., the3rd peak, the Pb-Bi spike and the Th, U abundances.The specific calculations employed here assume a

weighted range of neutron number densities (from 1023

to 1030 cm−3) and are designed to reproduce thetotal r-process isotopic S.S. abundance pattern (seeKratz et al. 2007a for more details). The best agree-ment with this abundance pattern, as well as with theelemental abundances in metal-poor halo stars, is ob-tained by employing the nuclear mass predictions—manyof the nuclei involved in this process are too short-lived to be experimentally measured at this time—ofthe ETFSI-Q nuclear mass model. We employed the β-decay properties from quasi-particle random-phase ap-proximation calculations for the Gamow-Teller transi-tions (see Moller & Randrup 1990; Moller et al. 1997)and the first-forbidden strength contributions from theGross theory of β-decay (Moller et al. 2003). We findthat a small number of individual neutron density com-ponents (in the current calculations, 15 components) isnecessary to fit the predicted r-process abundances to theS.S. values, and we also assume a varying r-process pathrelated to contour lines of constant neutron separationenergies in the range of 4–2 MeV. (This path is deter-mined by the variations of the neutron number densityand the temperature).The nuclear data for these calculations have been im-

proved by incorporating recent experimental results andimproved theoretical predictions. Analyses of the differ-ences between measured and predicted nuclear parame-ters (e.g., β-decay properties) indicate considerable im-provements over earlier attempts. This gives us confi-dence in the reliability of our nuclear physics input tothe r-process calculations of the heavy element region be-tween the rare earth elements, via the 3rd-peak elements(Os, Ir, Pt), the Pb and Bi isotopes, and up to Th and U.In addition the excellent agreement between these calcu-lations and the S.S. isotopic and elemental abundancessuggests that this approach can reproduce the astrophys-ical and nuclear conditions appropriate for the r-processdespite not knowing the astrophysical site for this pro-

Lead and Thorium from the r-process in Metal-Poor Stars 9

Fig. 6.— Comparison of the Hf/Eu and Ir/Eu ratios in our sample. Symbols are the same as in Figure 4. Downward facing trianglesrepresent upper limits.

cess (Kratz et al. 2007a). For this paper the theoreticalpredictions have been normalized to the r-process com-ponent (95%) of the S.S. Eu abundance (Lodders 2003).The predicted Pb abundance shown in Figure 9 is

broken into four components corresponding to its originwithin the r-process. “Direct production” refers to Pbthat is produced in the r-process as nuclei with A = 206,207, or 208, each of which will β− decay directly to oneof the stable Pb isotopes. “Indirect production” refersto Pb that is produced from the α and β decay of nucleiwith 210 ≤ A ≤ 231 and A = 234 shortly after the ter-mination of the r-process (i.e., within a few ×106 years).“Th and U decay” refers to Pb that originates from thedecay of nuclei that were produced in the r-process withA = 232, 235, or 238, which quickly β− decayed to 232Th,235U, and 238U, which have since decayed to the stablePb isotopes. In Figure 9 we show the Pb abundanceafter 13 Gyr. “Transuranic decay” refers to Pb that isproduced from the decay of 232Th, 235U, and 238U, butnow considering the fractions of the abundances of thesethree isotopes that were formed in the r-process as nu-clei with A = 236 and A ≥ 239, which followed α and βdecay chains to the long-lived Th and U isotopes.The abundance predictions for Pb, Th, and U, com-

puted at time “zero” after all of the α and β decaysare complete (∼ 107 years after the r-process event) are

listed in Table 8. According to these predictions, themajority of the r-process Pb abundance derives from in-direct production of short-lived nuclei between Pb andTh (82% at time “zero”). Only a small amount of thepresent-day Pb originated from the decay of the long-lived Th and U isotopes. The predicted total Pb/Euratio, based upon our waiting-point r-process calcula-tions as described above, is in agreement with the derivedPb/Eu ratios in all of the stars with −2.2 < [Fe/H] <−1.4, independent of the amount of time since ther-process nucleosynthesis event. This implies that fis-sion losses from nuclei with 210 ≤ A ≤ 231 and 234 arenot significant, otherwise the predicted total Pb abun-dance would be noticeably lower. These results, basedupon detailed fits that reproduce the total S.S. isotopicabundance distribution, are supported by recent dynamiccalculations (see, e.g., Farouqi et al. 2009), which alsoindicate that the total amount of fission material in thismass range will be insignificant and not affect the Pb toTh and U region. An additional argument regarding thelack of any appreciable effects of fission comes from recentfission barrier height calculations of Moller (2009, privatecommunication), who finds that any fissioning isotopesin the mass region 206 < A < 232–238—if they existat all—lie well beyond the r-process path in our modelpredictions (i.e., in isotopes more neutron-rich than the

10 Roederer et al.

Fig. 7.— Comparison of the Pb/Eu and Th/Eu ratios in our sample. Symbols are the same as in Figures 4 and 6.

r-process path and closer to the neutron drip line). Anyfissioning isotopes would not normally contribute to theabundances in the Pb region (even if they did not fission).Our results also indicate that any fissioning nuclei in themass region 232–238 < A < 250–256 make no significantcontributions to the stable heavy isotopes.We emphasize here that our approach has been to fit

the observed S.S. isotopic (and elemental) abundance dis-tribution with the superposition of the neutron numberdensity exposures that occur in the r-process. These cal-culations are further strengthened by utilizing the mostcomprehensive set of nuclear data currently available—including experimentally measured and theoretically de-termined (published and unpublished ETFSI-Q) masses,β-decay half-lives, and n-capture rates. Our normal pro-cedure has been to globally fit the S.S. abundances, par-ticularly normalizing the fit to the 2nd and 3rd processpeaks. In addition to employing these global fits wehave also fine-tuned our theoretical calculations to repro-duce the complete S.S. 3rd abundance peak in the massrange of Os (A = 186) to Hg (A = 204). It has longbeen argued that the production of the radioactive ele-ments extending upwards from Pb to Th and U and theheaviest stable elements in the 3rd r-process peak arecorrelated (Cowan et al. 1991; Thielemann et al. 1993;Cowan et al. 1999; Kratz et al. 2004, 2007a). Only the

correct reproduction of the A = 195 abundance peakwith its N = 126 (nuclear) bottle-neck behavior willguarantee that the extrapolated abundances into the ex-perimentally (completely) unknown r-process trans-Pbregion, including exotic isotopes about 30–40 units awayfrom the known “valley of the isobaric mass parabolae,”will be reliable.The excellent agreement between the r-process Pb

abundances in stars with −2.2 < [Fe/H] < −1.4 andour predictions is encouraging. Assuming the stellar Pbabundances are not seriously in error, we currently lacka complete, self-consistent understanding of r-processnucleosynthesis and enrichment for all low metallicitystars. Both CS 31082–001 and HE 1523–0901 have simi-lar metallicities and levels of r-process enrichment. TheirU/Th ratios are similar, and their Pb abundances maybe similar to one another (yet different from more metal-rich r-process-enriched stars). Their Th/Eu ratios aredifferent from one another, yet the Th/Eu ratio inHE 1523–0901 agrees with most other metal-poor stars,while CS 31082–001 does not. (See further discussionin Frebel & Kratz 2009.) A still larger set of r-process-enriched stars with [Fe/H] < −2.2 will be necessary tocharacterize the abundance patterns of Pb, Th, and U atlow metallicity, and until that time (at least) our under-standing of the nucleosynthesis of the heaviest products

Lead and Thorium from the r-process in Metal-Poor Stars 11

Fig. 8.— Comparison of the mean ratios for stars exhibiting apure r-process signature (log ǫ (La/Eu) < +0.25; red diamonds)with the four “actinide boost” stars (green circles). The predictedS.S. r-process“residual” abundance pattern is shown for reference.The abundances are normalized at Eu. Different decay ages areindicated by the dotted (t = 0 Gyr), dashed (t = 4.6 Gyr), andsolid (t = 13.0 Gyr) lines. Any deviations in the abundances of the“actinide boost” stars from the “standard” r-process stars clearlyoccur only after the 3rd r-process peak.

Fig. 9.— Comparison of r-process-only stellar Pb/Eu ratioswith Pb abundance predictions. Symbols are the same as in Fig-ure 6. The Solar r-process ratio is indicated by the ⊙. PredictedPb/Eu ratios are indicated by horizontal lines, representing con-tributions to the present-day Pb abundance via direct productionof the 206,207,208Pb isotopes; indirect Pb production via α and βdecays from nuclei with 210 ≤ A ≤ 231 and A =234 shortly afterthe termination of the r-process event; 13 Gyr of decay from nu-clei that formed along the isobars of 232Th, 235U, and 238U; and13 Gyr of decay from the fraction of these three isotopes producedindirectly from the α and β decay of transuranic nuclei shortly af-ter the r-process event has shut off. The excellent agreement of thepredicted total Pb/Eu ratio with the stellar values (with the ex-ceptions of CS 31082–001 and HE 1523–0901) implies that fissionlosses from the region between Pb and Th are not significant.

of the r-process will remain incomplete.

6. STELLAR AGES

When deriving stellar ages through nuclear chronom-etry, the measured ages reflect the ages of the actinidesin these stars but not necessarily the ages of the starsthemselves. In addition, a few r-process events may haveseeded the ISM from which these stars formed. Presum-ably the stars formed shortly after the r-process materialwas created and so the derived age represents an upperlimit to—but also a realistic estimate of—the stellar age.U/r ratios (where r denotes a stable element produced

in the same r-process event as the U) have the strongestpredictive power (0.068 dex per Gyr), followed by U/Th(0.046 dex per Gyr), and Th/r (0.021 dex per Gyr).These rates are only governed by the nuclear half-livesof the radioactive isotopes, which have been measured toexquisite precision as far as stellar nuclear chronometryis concerned. Strictly speaking, the Pb/r (or U/Pb, orTh/Pb) ratio is also a chronometer, since the Pb abun-dance increases with time as the Th and U decay, but itloses sensitivity as time passes and most of the Th and Unuclei decay. According to our model, 90% of the totalincrease in the Pb abundance has already occurred in thefirst 5 Gyr after the r-process event, and the total (i.e.,t = ∞) Pb abundance is only 0.10 dex higher than itstime “zero” abundance.Stellar abundance ratios are rarely known to better

than 0.05–0.10 dex, limiting relative age determinationsfrom one star to another using a single Th/r pair to aprecision of ∼ 3–5 Gyr. To some degree this can bemitigated by employing multiple Th/r pairs, but uncer-tainties in the production ratios and uncertainties arisingfrom systematic effects (e.g., determination of effectivetemperature) will limit the measurement of an absoluteage to no better than ∼ 2–3 Gyr (see, e.g., Sneden et al.2000 or Frebel et al. 2007).Figure 10 illustrates these points by showing the

ages derived for our r-process-only sample from Th/Eu,Th/Hf, Pb/Th, and U/Th chronometers.7 Our results

7 The element Hf has recently been suggested as a promis-ing, stable r-process reference element. Lundqvist et al. (2006)and Lawler et al. (2007) have made new laboratory evaluations ofHf ii transition probabilities, and there are indications that Hfmay be formed at similar neutron densities to the 3rd r-processpeak elements Os, Ir, and Pt (Kratz et al. 2007a). These elementsare difficult to reliably measure in stellar spectra obtained fromground-based facilities (except for Ir) and can only be observedin their neutral states. In other words, Hf may be the heavieststable singly-ionized element whose production is closely linkedto that of the actinides and may be reliably and easily measuredin metal-poor stars. Our model predicts an r-process Hf abun-dance of 0.0436 (26% of the S.S. value of 0.1699 for A = 174–180,Lodders 2003). The isotopic and elemental Hf r-process abun-dance predictions by our model are consistently smaller than thoseinferred from the S.S. r-process residual method. We caution thatknowledge of the nuclear structure of atoms is incomplete for iso-topes near A ∼ 180, in the transition from the deformed rare earthisotopes to the spherical N = 126 magic shell. The offset be-tween the scaled S.S. and stellar r-process Hf abundances noticedby Lawler et al. (2007) and Sneden et al. (2009) suggests that alarger fraction of the S.S. Hf abundance might be attributed tothe r-process. That study adopted the S.S. Hf r-process fraction(44%) from Arlandini et al. (1999), implying that the S.S. r-processfraction is even higher than 44%. If we reverse the problem andassume that the stellar Hf/Eu r-process ratio from Sneden et al.(2009) should match the S.S. Hf/Eu r-process ratio, we estimatethe r-process fraction of Hf in the S.S. is ≈ 70±10%.

12 Roederer et al.

are consistent with the assertion that r-process enrich-ment began extremely early, within the first few Gyr af-ter the big bang. The horizontal lines indicating agesare determined based on the production ratios given inKratz et al. (2007a). Using a different set of produc-tion ratios (e.g., Cowan et al. 1999, Schatz et al. 2002, orthose predicted by our stellar sample) would only changethe absolute scale of these age determinations by smallamounts (∼ 2–4 Gyr; see Frebel et al. 2007).This is the largest sample of Th/Eu ratios yet com-

piled for metal-poor halo stars. The chronometer pairsin an individual star have limited ability to predict itsage, but the combined measurements from an ensembleof stars hold greater promise. The majority of our samplesuggests an old population (the exception being the fourstars with an actinide boost), and no trends with metal-licity are apparent. This is in agreement with modelpredictions for the age-metallicity relationship in metal-poor halo stars, where a rapid increase is found in themean metallicity to [Fe/H] ∼ −2.0 within the first 2 Gyr(e.g., Argast et al. 2000). If we divide the sample intotwo groups of stars—those with an actinide boost andthose without—and assume a single age for each group,we can derive reasonable estimates for the age of ther-process-only standard stars, as shown in Table 9. As-suming that the observed stellar ratios are independent(which they clearly are not since all rely on Th),8 Wederive an age for the ensemble of standard r-process-only stars of 15.2±2.1 (σ=4.6) Gyr. Two ratios predictages significantly greater than the age of the Universe,Th/La (20.4±4.2 Gyr) and Th/Hf (19.7±4.7 Gyr). If weadopt the higher S.S. r-process “residual” abundances forLa and Hf (e.g., Sneden et al. 2008, with updates fromGallino), the Th/La and Th/Hf ages decrease by 13 and3 Gyr, respectively.The four stars with an actinide boost reflect a very

different age, 1.3±2.1 (σ=4.6) Gyr. In CS 31082–001,where abundances of both Th and U have been derived(Cayrel et al. 2001; Hill et al. 2002), the U/Th ratio im-plies a reasonable age of 14.2±2.4 (±sys) Gyr (Hill et al.2002). If we use the 238U/232Th production ratio com-puted from the data in Table 8 and assume a 20% un-certainty in this value, we derive an age of 15.1±4.3 Gyrin CS 31082–001. The low Pb abundance and high Th/rratios in this star cannot be accounted for simply byassuming a very young age; increasing the current Pbabundance in this star by the maximum expected fromthe complete decay of the actinides (≈0.10 dex) wouldstill leave it about 0.4 dex lower than the mean of themore metal-rich r-process enriched stars. U has also beendetected in two r-process enhanced stars that do not havean actinide boost, BD+17 3248 (Cowan et al. 2002) andHE 1523–0901 (Frebel et al. 2007), and the U/Th agesin these two stars are consistent with the U/Th age in

8 A change in the Th abundance by ±0.05 dex will uniformlychange the derived age by ∓2.3 Gyr. The uncertainty on the agederived from each chronometer pair is computed by combining inquadrature the error of the mean stellar ratio, an assumed 20%uncertainty in the production ratio, and the uncertainty in theabundances arising from uncertainties in the stellar parameters.For this final source of error, we consider the uncertainties derivedby Frebel et al. (2007) for HE 1523–0901 to be representative of theuncertainties for an individual star in our sample, and we reduceeach of their uncertainties by the amount expected when increasingthe sample size from one star to several.

CS 31082–001 (see Figure 10).

6.1. Comparison to Globular Cluster and dSph AgeEstimates from Other Methods

Several stars in our sample are located in globularclusters or dSph systems, whose ages and star forma-tion histories can be estimated based on other methods.M15 is an old (13.2±1.5 Gyr, McNamara et al. 2004),metal-poor ([Fe/H] ∼ −2.2; Gratton et al. 2004) globu-lar cluster. Sneden et al. (2000) derived abundances ofseveral n-capture elements in three giants in this clus-ter. The mean Th/Eu ratio for these three stars im-plies an age of 12.0±3.7 Gyr using our production ratio.M92 is also an old (14.8±3 Gyr, Carretta et al. 2000;14.2±1.2 Gyr, Paust et al. 2007), metal-poor ([Fe/H] ∼−2.3; Gratton et al. 2004) globular cluster. Using theTh/Eu ratio derived by Johnson & Bolte (2001) for onestar in this cluster, we find and age of 10.6±4.7 Gyr.Aoki et al. (2007) derived the Th/Eu ratio in one star inUMi, which implies an age of 12.0±6.5 Gyr. This resultis consistent with the earlier finding that this system ex-perienced only one significant episode of star formationat early times (& 11 Gyr ago; Dolphin 2002). In allthree cases, the ages for these systems derived from nu-clear chronometry and other independent methods agreewithin the uncertainties.

7. IMPLICATIONS FOR NUCLEOSYNTHESIS IN THEEARLY GALAXY

One of the more remarkable aspects of the abundanceratios for r-process-only stars in Figures 4, 6, and 7 is thewide range of Eu and Fe abundances covered by thesecorrelations. The r/Eu ratios are constant over a verywide range of absolute r-process enrichment, roughly2.4 dex or a factor of 250. All of the [r/Eu] abun-dance ratios in these figures (with the noted exceptionof the [Pb/Eu] and [Th/Eu] ratios in the stars with anactinide boost) are unchanged over the metallicity range−3.1 ≤ [Fe/H]≤ −1.4. For these stars, the [Eu/Fe] ratiosare always supersolar, but they vary widely, from +0.31(M92 VII-18) to +1.82 (HE 1523–0901). This wide dis-persion in n-capture abundances at low metallicities hasbeen previously noted by many investigators, includ-ing Gilroy et al. (1988) and Burris et al. (2000). Severalstars with +0.3 < [Eu/Fe] < +0.5 exhibit an r-process-only signature, which reveals—as might be expected—that a pure r-process pattern can be found even insmall amounts of r-process enrichment. Furthermore,this same pattern is observed in stars in the Galactichalo, several globular clusters, and one dSph system.Taken together, these facts are strong evidence for theuniversal nature of the main r-process for species withZ ≥ 56 (as constrained by observations), since starswith log ǫ (Eu) ∼ +0.4 dex certainly are comprised ofthe remnants of many more supernovae than stars withlog ǫ (Eu) ∼ −2.0.From an analysis of the La/Eu ratio in a sample of

159 stars with −3.0 < [Fe/H] < +0.3, Simmerer et al.(2004) found stars with [Fe/H] = −2.6 exhibiting sig-natures of the s-process, while other stars as metal-rich as [Fe/H] = −1.0 showed little evidence of anys-process material. By the definition we adopt in Sec-tion 4, log ǫ (La/Eu)r < +0.25, processes other thanthe r-process (e.g., the so-called weak r-process) must

Lead and Thorium from the r-process in Metal-Poor Stars 13

Fig. 10.— Four nuclear chronometer pairs. Only stars with a pure r-process signature are shown. Symbols are the same as in Figure 6.The horizontal lines indicate the ratios expected if a sample of material had a given age, assuming the nucleosynthesis predictions ofKratz et al. (2007a). The vertical scales cover the same number of decades on each panel, though the ranges differ, to illustrate the relativemeasurement precision. The shortcomings of the r-process Hf predictions by our model are evident in the poor match between the Th/Hfpredictions and observations shown in the lower left panel.

be responsible for some n-capture material at even lowermetallicity. Simmerer et al. (2004) also found incompletemixing of both r-process and s-process material in starswith [Fe/H] > −1.0. We find that the gas from whichthese stars formed was inhomogeneous at metallicities ashigh as [Fe/H] = −1.4, the metal-rich limit of our sam-ple; here, several stars show no evidence of any contribu-tions from the s-process. Furthermore, Roederer (2009)found no preferred kinematic signature for stars withpure r-process or s-process enrichment patterns, withthese patterns extending over a wide metallicity range of−3.0 < [Fe/H] < −0.4. This reinforces the notion thatn-capture enrichment at low metallicity is likely a verylocalized phenomenon that results in a large distributionof n-capture abundances.The range of absolute r-process enrichment noted

above includes 18 stars in M15. (Three stars fromSneden et al. 2000 are shown in Figure 4; additional starsfrom Sneden et al. 1997 and Otsuki et al. 2006, who didnot report Th abundances, are not shown.) These 18stars in a single globular cluster show no change in theirBa/Eu or La/Eu ratios despite a change in the absoluteEu enrichment level by 0.9 dex, a factor of ≈ 8; their

Fe abundances differ by less than 0.2 dex, a factor of≈ 1.5. Sneden et al. (1997) found no correlation betweenthese n-capture enrichment patterns in M15 and the sig-natures of deep mixing commonly observed in globularcluster stars, indicating a primordial origin. This enrich-ment pattern resembles that of the Galactic halo, but itis not obvious why other clusters enriched by r-processmaterial (with or without significant contributions fromthe s-process) fail to show similar star-to-star dispersion.

8. CONCLUSIONS

We have identified a sample of 27 stars with−3.1 < [Fe/H] < −1.4 that have been enriched by ther-process and show no evidence of s-process enrichment.We confront r-process nucleosynthesis predictions for Pband Th with measurements (or upper limits) in our stel-lar sample. We use these very heavy isotopes locatednear the termination points of s- and r-process nucle-osynthesis to better understand the physical nature ofthe r-process and the onset of nucleosynthesis in the earlyGalaxy. Our major results can be summarized as follows.Stars with log ǫ (La/Eu) < +0.25 (our “r-process-

only” sample, where more than ≈ 99% of the n-capturematerial originated in the r-process) show no evolu-

14 Roederer et al.

tion in their Pb/La ratio over the metallicity range−2.2 < [Fe/H] < −1.4. In contrast, stars withlog ǫ (La/Eu) ≥ +0.25 (those with just a dusting ofs-process material on top of r-process enrichment) show asignificant increase in Pb/La with decreasing metallicity.This emphasizes the effect of the higher neutron-to-seedratio that occurs in low metallicity s-process environ-ments and overproduces nuclei at the termination of thes-process chain relative to lighter s-process nuclei. Thiseffect is noticeable in stars where only ≈ 2.0% of then-capture material originated in the s-process.All stars in our r-process-only sample have constant

abundance ratios among elements surveyed in the rareearth domain and the 3rd r-process peak (La, Eu, Er,Hf, and Ir), and these abundance ratios are equivalentto the scaled S.S. r-process distribution. These ratiosare identical in stars with a so-called actinide boost andstars without. For stars with an actinide boost, our ob-servations demonstrate that any nucleosynthetic devia-tions from the main r-process affect—at most—only theelements beyond the 3rd r-process peak (Pb, Th, and U).We find very good agreement between the Pb abun-

dances in our r-process-only stars and the Pb abundancespredicted by our classical “waiting-point” r-processmodel. In these computations a superposition of15 weighted neutron-density components in the range23 ≤ log nn ≤ 30 is used to successfully reproduceboth the S.S. isotopic distribution and the heavy ele-ment abundance pattern between Ba and U in the low-metallicity halo stars. Our calculations indicate that fis-sion losses are negligible for nuclei along the r-processpath between Pb and the radioactive isotopes of Th andU. In light of this agreement, we currently have no vi-able theoretical explanation for the low Pb abundance inCS 31082–001.With the exception of the Pb and Th in stars with

an actinide boost, the r/Eu ratios in our r-process-only sample are constant over a wide range of metal-licity (−3.1 < [Fe/H] < −1.4) and r-process enrichment(−2.0 < log ǫ (Eu) < +0.4 or +0.3 < [Eu/Fe] < +1.8).This pattern is observed in field stars, several globular

clusters, and one dSph system. As multiple supernovaewill have contributed to the material in stars at the high-est metallicities and/or r-process enrichments surveyed,we regard this as strong evidence for the universal natureof the r-process.We have derived an age of 15.2±2.1 Gyr (σ = 4.6 Gyr)

from several Th/r chronometer pairs for an ensemble of16 stars. This is the largest set of Th/r ratios yet com-piled for metal-poor halo stars. Excluding the four starswith an actinide boost, there is no relationship betweenage and metallicity over the range−3.1 < [Fe/H] < −1.4.While each stellar chronometer pair ratio argues for asingle common age for the ensemble of stars, the obser-vations call for further refinement of some productionratios, and the outlook for improving the precision ofthe age measurement for a single star in the field is nothopeful. For two globular clusters and one dSph system,we do find good agreement (within uncertainties ∼ 3.5–5 Gyr) between ages derived from the Th/Eu nuclearchronometer and ages derived from other methods.

We thank Satoshi Honda and David Lai for sendingunpublished upper limit measurements and David Yongfor providing additional abundance results. We appre-ciate the helpful suggestions of the referee. I.U.R. alsothanks Jim Truran for an invigorating discussion on thistopic. This research has made use of the NASA As-trophysics Data System (ADS), NIST Atomic SpectraDatabase, and the SIMBAD database, operated at CDS,Strasbourg, France. A.F. acknowledges support throughthe W.J. McDonald Fellowship of the McDonald Obser-vatory. Funding for this project has also been generouslyprovided by the Deutsche Forschungsgemeinschaft (con-tract KR 806/13-1), the Helmholtz Gemeinschaft (grantVH-VI-061), and the U. S. National Science Foundation(grants AST 07-07447 to J.J.C. and AST 06-07708 toC.S.).Facilities: HST (STIS) Keck:I (HIRES) VLT:Kueyen

(UVES)

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16 Roederer et al.

TABLE 1Adopted Atmospheric Parameters and Spectral Coverage

Star Teff log g vt [Fe/H] Reference HST Keck VLT(K) (km s−1) STIS HIRES UVES

BD+17 3248 5200 1.80 1.90 −2.08 1 Y Y NCS 22892–052 4800 1.50 1.95 −3.10 2 Y Y NCS 29491–069 5300 2.80 1.60 −2.60 3 N N YHD 6268 4685 1.50 2.00 −2.42 4 Y Y NHD 74462 4700 2.00 1.90 −1.52 5 N Y NHD 108317 5234 2.68 2.00 −2.18 5 N Y NHD 115444 4720 1.75 2.00 −2.90 5 Y Y NHD 122563 4570 1.35 2.90 −2.72 5 Y Y NHD 122956 4510 1.55 1.60 −1.95 5 Y Y NHD 126587 4795 1.95 2.00 −2.93 4 Y Y NHD 175305 5040 2.85 2.00 −1.48 4 Y Y NHD 186478 4600 1.45 2.00 −2.56 5 Y Y NHD 204543 4672 1.49 2.00 −1.87 5 N Y NHE 1219–0312 5060 2.30 1.60 −2.97 3 N N Y

References. — (1) Cowan et al. 2002; (2) Sneden et al. 2003; (3) Hayek et al. 2009;(4) Cowan et al. 2005; (5) Simmerer et al. 2004

LeadandThoriu

mfro

mther-p

rocess

inMeta

l-PoorStars

17

TABLE 2Abundances Dervied from Individual Transitions (I)

λ (A) species E.P. (eV) log(gf) Ref. BD+17 3248 CS 22892–052 CS 29491–069 HD 6268 HD 74462 HD 108317 HD 115444

3794.77 La ii 0.24 +0.21 1 −0.58±0.15 −0.91±0.10 −0.75±0.15 −1.08±0.15 · · · −0.96±0.20 −1.38±0.103988.51 La ii 0.40 +0.21 1 −0.57±0.10 −0.88±0.10 −0.80±0.15 −1.08±0.10 −0.24±0.10 −1.04±0.15 −1.48±0.103995.74 La ii 0.17 −0.06 1 −0.54±0.10 −0.86±0.10 −0.73±0.15 −1.05±0.10 −0.24±0.10 −0.97±0.15 −1.43±0.104086.71 La ii 0.00 −0.07 1 −0.52±0.10 −0.83±0.10 −0.70±0.10 −0.97±0.10 −0.30±0.20 −0.96±0.20 −1.36±0.104123.22 La ii 0.32 +0.13 1 −0.59±0.15 −0.90±0.15 −0.79±0.10 −1.12±0.15 −0.30±0.20 −1.07±0.15 −1.48±0.153724.93 Eu ii 0.00 −0.09 2 −0.75±0.20 −0.94±0.20 −0.98±0.20 −1.31±0.15 −0.49±0.25 −1.30±0.15 −1.61±0.153819.67 Eu ii 0.00 +0.51 2 −0.81±0.15 −0.94±0.15 −0.97±0.10 −1.41±0.15 · · · −1.35±0.15 −1.68±0.103907.11 Eu ii 0.21 +0.17 2 −0.81±0.10 −0.96±0.10 −0.95±0.10 −1.39±0.10 −0.53±0.25 −1.31±0.10 −1.64±0.103930.50 Eu ii 0.21 +0.27 2 −0.77±0.20 −0.97±0.20 −0.97±0.20 −1.30±0.20 · · · −1.28±0.20 −1.63±0.203971.97 Eu ii 0.21 +0.27 2 −0.81±0.20 −0.99±0.20 −0.97±0.20 −1.40±0.20 · · · −1.36±0.20 −1.64±0.204129.72 Eu ii 0.00 +0.22 2 −0.77±0.10 −0.98±0.10 −0.95±0.10 −1.39±0.10 −0.49±0.15 −1.32±0.10 −1.63±0.104205.04 Eu ii 0.00 +0.21 2 −0.77±0.15 −0.97±0.20 −0.97±0.15 −1.36±0.15 −0.52±0.20 · · · −1.62±0.104435.58 Eu ii 0.21 −0.11 2 −0.73±0.20 · · · −0.96±0.15 · · · −0.50±0.20 · · · −1.66±0.153230.58 Er ii 0.06 +0.24 3 −0.43±0.20 −0.62±0.30 · · · −0.98±0.15 −0.16±0.30 −0.92±0.20 −1.29±0.203312.43 Er ii 0.06 −0.03 3 −0.32±0.20 −0.56±0.25 · · · −0.85±0.20 · · · −0.76±0.25 −1.20±0.153729.52 Er ii 0.00 −0.59 3 −0.27±0.15 −0.44±0.15 −0.54±0.15 −0.85±0.15 −0.05±0.15 −0.82±0.15 −1.14±0.153830.48 Er ii 0.00 −0.22 3 −0.36±0.15 −0.51±0.15 −0.50±0.15 −0.96±0.15 −0.24±0.20 −0.77±0.15 −1.23±0.153896.23 Er ii 0.06 −0.12 3 −0.34±0.15 −0.50±0.15 −0.55±0.15 −1.01±0.20 −0.26±0.30 −0.92±0.15 −1.29±0.153906.31 Er ii 0.00 +0.12 3 −0.38±0.20 −0.52±0.20 −0.47±0.20 · · · · · · −0.88±0.20 −1.21±0.203193.53 Hf ii 0.38 −0.89 4 −0.85±0.25 · · · · · · · · · −0.29±0.30 < −0.80 · · ·3505.22 Hf ii 1.04 −0.14 4 −0.75±0.20 −1.07±0.30 · · · −1.09±0.25 · · · −1.00±0.20 −1.52±0.303918.09 Hf ii 0.45 −1.14 4 · · · −0.81±0.20 −0.57±0.25 −1.04±0.25 · · · · · · · · ·4093.15 Hf ii 0.45 −1.15 4 −0.68±0.15 −0.89±0.15 −0.81±0.30 −1.24±0.20 −0.39±0.10 −1.00±0.30 −1.47±0.303513.65 Ir i 0.00 −1.21 5 +0.19±0.25 −0.07±0.20 +0.21±0.30 −0.41±0.20 +0.59±0.25 −0.22±0.25 −0.49±0.253800.12 Ir i 0.00 −1.44 5 +0.08±0.20 −0.14±0.20 +0.06±0.30 −0.57±0.20 +0.37±0.15 −0.27±0.20 −0.86±0.252833.03 Pb i 0.00 −0.50 6, 7 < +0.27 < −0.15 · · · · · · · · · · · · < −0.453683.46 Pb i 0.97 −0.54 6, 7 < +0.42 < +0.20 < +0.35 < +0.18 +0.53±0.20 +0.17±0.30 < +0.054057.81 Pb i 1.32 −0.22 6, 7 < +0.72 < +0.35 < +0.50 < +0.08 < +0.48 < +0.37 < −0.303539.59 Th ii 0.00 −0.54 8 −1.31±0.20 −1.55±0.25 · · · < −1.58 −1.17±0.30 < −1.44 < −1.864019.13 Th ii 0.00 −0.23 8 −1.27±0.15 −1.68±0.20 −1.46±0.25 −1.78±0.15 −0.88±0.20 −1.84±0.20 −2.08±0.154086.52 Th ii 0.00 −0.93 8 −1.14±0.30 −1.48±0.30 < −1.06 < −1.33 −0.90±0.20 < −1.19 < −1.714094.75 Th ii 0.00 −0.88 8 −1.19±0.30 −1.60±0.30 < −1.16 < −1.43 < −0.88 < −1.04 < −1.81

References. — (1) Lawler et al. 2001a; (2) Lawler et al. 2001b; (3) Lawler et al. 2008; (4) Lawler et al. 2007; (5) Ivarsson et al. 2003, with updatesas noted in the Appendix of Cowan et al. 2005; (6) Wood & Andrew 1968; (7) Biemont et al. 2000; (8) Nilsson et al. 2002

18

Roederer

etal.

TABLE 3Abundances Dervied from Individual Transitions (II)

λ (A) species E.P. (eV) log(gf) Ref. HD 122563 HD 122956 HD 126587 HD 175305 HD 186478 HD 204543 HE 1219–0312

3794.77 La ii 0.24 +0.21 1 · · · −0.78±0.25 −1.75±0.15 −0.16±0.20 −1.34±0.20 · · · −0.75±0.203988.51 La ii 0.40 +0.21 1 −2.50±0.30 −0.63±0.10 −1.83±0.25 −0.14±0.10 −1.33±0.10 −0.63±0.10 −0.80±0.203995.74 La ii 0.17 −0.06 1 −2.32±0.30 −0.63±0.10 −1.82±0.15 −0.12±0.10 −1.34±0.10 −0.62±0.10 −0.73±0.154086.71 La ii 0.00 −0.07 1 −2.34±0.20 −0.66±0.10 −1.59±0.10 −0.17±0.10 −1.28±0.10 −0.63±0.10 −0.72±0.104123.22 La ii 0.32 +0.13 1 −2.48±0.25 −0.58±0.20 −1.81±0.20 −0.13±0.15 −1.35±0.15 −0.65±0.15 −0.88±0.203724.93 Eu ii 0.00 −0.09 2 −2.77±0.25 −0.85±0.25 −1.90±0.20 −0.31±0.15 −1.49±0.15 −1.03±0.15 −0.97±0.203819.67 Eu ii 0.00 +0.51 2 −2.91±0.30 · · · −2.04±0.15 · · · −1.56±0.10 · · · −1.00±0.103907.11 Eu ii 0.21 +0.17 2 −2.74±0.25 −0.98±0.15 −1.97±0.15 −0.38±0.20 −1.55±0.10 −1.14±0.15 −0.99±0.103930.50 Eu ii 0.21 +0.27 2 < −2.01 · · · −1.96±0.20 · · · · · · · · · −1.03±0.203971.97 Eu ii 0.21 +0.27 2 < −2.51 −1.03±0.20 −1.95±0.20 · · · −1.53±0.20 · · · −1.04±0.204129.72 Eu ii 0.00 +0.22 2 −2.71±0.20 −0.93±0.15 −1.95±0.10 −0.37±0.10 −1.51±0.10 −1.04±0.15 −0.97±0.104205.04 Eu ii 0.00 +0.21 2 −2.68±0.30 −0.93±0.20 −2.02±0.20 −0.38±0.25 −1.53±0.15 −1.05±0.15 −0.95±0.154435.58 Eu ii 0.21 −0.11 2 < −2.21 −0.88±0.20 < −1.82 −0.40±0.25 −1.55±0.20 −0.99±0.15 −0.91±0.203230.58 Er ii 0.06 +0.24 3 −2.38±0.30 −0.31±0.20 −1.47±0.25 −0.06±0.20 −1.10±0.20 −0.68±0.25 · · ·3312.43 Er ii 0.06 −0.03 3 < −2.11 −0.24±0.20 −1.47±0.30 −0.02±0.30 −1.06±0.25 −0.48±0.25 · · ·3729.52 Er ii 0.00 −0.59 3 < −1.76 −0.45±0.15 −1.43±0.20 +0.03±0.15 −1.03±0.15 −0.63±0.15 −0.49±0.153830.48 Er ii 0.00 −0.22 3 −2.23±0.30 −0.46±0.15 −1.46±0.15 −0.03±0.15 −1.13±0.15 −0.75±0.20 −0.52±0.153896.23 Er ii 0.06 −0.12 3 < −2.16 −0.49±0.25 −1.56±0.20 +0.00±0.20 −1.16±0.20 −0.58±0.25 −0.50±0.153906.31 Er ii 0.00 +0.12 3 · · · · · · −1.46±0.20 · · · −1.03±0.25 · · · −0.45±0.203193.53 Hf ii 0.38 −0.89 4 · · · · · · · · · −0.25±0.30 · · · −0.89±0.30 · · ·3505.22 Hf ii 1.04 −0.14 4 < −1.64 −0.69±0.25 < −1.40 −0.22±0.20 −1.23±0.25 −0.77±0.20 · · ·3918.09 Hf ii 0.45 −1.14 4 · · · −0.48±0.15 < −1.05 · · · · · · · · · < −0.544093.15 Hf ii 0.45 −1.15 4 < −1.89 −0.77±0.15 −1.65±0.30 −0.05±0.15 −1.40±0.15 −0.83±0.15 < −0.393513.65 Ir i 0.00 −1.21 5 < −0.51 −0.08±0.20 < −0.47 +0.40±0.20 −0.62±0.25 −0.13±0.25 · · ·3800.12 Ir i 0.00 −1.44 5 < −1.16 −0.09±0.20 −1.05±0.25 +0.41±0.20 −0.79±0.20 −0.27±0.20 −0.19±0.252833.03 Pb i 0.00 −0.50 6, 7 < −0.42 −0.15±0.20 < −0.38 +0.07±0.30 · · · · · · · · ·3683.46 Pb i 0.97 −0.54 6, 7 < −0.02 +0.00±0.20 < +0.07 +0.22±0.30 < −0.01 +0.05±0.25 < +0.634057.81 Pb i 1.32 −0.22 6, 7 < −0.27 < +0.05 < −0.28 < +0.42 < −0.26 < +0.13 < +0.533539.59 Th ii 0.00 −0.54 8 < −1.88 −1.76±0.30 < −1.84 −0.84±0.30 < −1.92 −1.96±0.30 < −0.484019.13 Th ii 0.00 −0.23 8 < −2.43 −1.38±0.20 < −2.39 −0.74±0.20 −2.19±0.25 −1.53±0.25 −1.25±0.254086.52 Th ii 0.00 −0.93 8 < −2.08 < −1.66 < −1.79 < −0.64 < −2.02 −1.71±0.30 −1.03±0.404094.75 Th ii 0.00 −0.88 8 < −1.93 < −1.41 < −1.79 −0.71±0.30 < −1.92 −1.60±0.30 < −0.88

References. — (1) Lawler et al. 2001a; (2) Lawler et al. 2001b; (3) Lawler et al. 2008; (4) Lawler et al. 2007; (5) Ivarsson et al. 2003, with updatesas noted in the Appendix of Cowan et al. 2005; (6) Wood & Andrew 1968; (7) Biemont et al. 2000; (8) Nilsson et al. 2002

LeadandThoriu

mfro

mther-p

rocess

inMeta

l-PoorStars

19

TABLE 4Final Elemental Abundances

Star [Fe/H] Ref. log ǫ (La) log ǫ (Eu) log ǫ (Er) log ǫ (Hf) log ǫ (Ir) log ǫ (Pb) log ǫ (Th) Ref. r-only?a

BD−18 5550 −3.05 1 −2.52±0.10 −2.81±0.20 −2.37±0.10 < −0.86 · · · · · · < −3.02 1, 2BD+01 2916 −1.93 3 −0.87±0.12 −1.22±0.14 · · · · · · · · · −0.20±0.19 · · · 3BD+04 2621 −2.52 1 −2.29±0.11 −2.63±0.20 −2.24±0.10 · · · · · · · · · < −3.04 1, 2BD+06 0648 −2.14 3 −0.95±0.12 −1.50±0.16 · · · · · · · · · < +0.00 · · · 3BD+08 2856 −2.12 1 −1.03±0.04 −1.16±0.04 −0.95±0.09 < +0.20 · · · · · · −1.70±0.10 1, 2 YBD+17 3248 −2.08 4 −0.55±0.05 −0.78±0.05 −0.34±0.07 −0.73±0.11 +0.12±0.16 < +0.27 −1.26±0.10 5 YBD+30 2611 −1.46 3 −0.27±0.12 −0.49±0.14 · · · · · · · · · +0.43±0.19 · · · 3 YCS 22892–052 −3.10 6 −0.87±0.05 −0.96±0.05 −0.50±0.07 −0.93±0.13 −0.10±0.14 < −0.15 −1.60±0.13 5 YCS 29491–069 −2.60 7 −0.75±0.05 −0.96±0.05 −0.52±0.08 −0.67±0.19 +0.13±0.21 < +0.35 −1.46±0.25 5 YCS 29497–004 −2.63 8 −0.38±0.15 −0.45±0.20 · · · · · · · · · · · · −0.96±0.15 8, 9 YCS 30306–132 −2.50 10 −0.78±0.06 −1.02±0.05 −0.62±0.15 · · · · · · < +0.50 −1.16±0.15 10 YCS 31078–018 −2.84 11 −1.00±0.22 −1.17±0.17 −0.99±0.15 · · · · · · < +0.25 −1.35±0.25 12 YCS 31082–001 −2.90 12 −0.62±0.04 −0.72±0.03 −0.30±0.04 −0.72±0.04 +0.22±0.20 −0.55±0.15 −0.98±0.05 12, 13, 14 YHD 3008 −1.98 3 −1.02±0.14 −1.30±0.16 · · · · · · · · · −0.35±0.22 · · · 3HD 6268 −2.42 15 −1.05±0.05 −1.37±0.05 −0.93±0.07 −1.14±0.13 −0.49±0.14 < +0.08 −1.78±0.15 5HD 29574 −1.86 3 −0.63±0.12 −0.90±0.14 · · · · · · · · · −0.15±0.19 · · · 3HD 74462 −1.52 16 −0.25±0.06 −0.50±0.09 −0.14±0.10 −0.38±0.09 +0.43±0.13 +0.53±0.20 −0.94±0.13 5HD 108317 −2.18 16 −1.01±0.07 −1.32±0.05 −0.85±0.07 −1.00±0.17 −0.25±0.16 +0.17±0.30 −1.84±0.20 5HD 108577 −2.38 1 −1.24±0.09 −1.48±0.12 −1.23±0.11 · · · · · · · · · −2.03±0.14 1, 2 YHD 115444 −2.90 16 −1.42±0.05 −1.64±0.04 −1.22±0.07 −1.50±0.21 −0.68±0.18 < −0.45 −2.08±0.15 5 YHD 122563 −2.72 16 −2.40±0.13 −2.75±0.11 −2.30±0.21 < −1.89 < −1.16 < −0.42 < −2.43 5HD 122956 −1.95 16 −0.64±0.05 −0.94±0.07 −0.40±0.08 −0.83±0.08 −0.08±0.14 −0.08±0.14 −1.50±0.17 5HD 126587 −2.93 15 −1.75±0.07 −1.97±0.06 −1.47±0.08 −1.65±0.30 −1.05±0.25 < −0.38 < −2.39 5 YHD 128279 −2.40 1 −1.51±0.16 −1.80±0.20 −1.41±0.10 · · · · · · · · · < −1.88 1, 2HD 141531 −1.66 3 −0.55±0.12 −0.87±0.14 · · · · · · · · · +0.05±0.19 · · · 3HD 175305 −1.48 15 −0.14±0.05 −0.36±0.07 −0.01±0.08 −0.13±0.11 +0.40±0.14 +0.25±0.21 −0.76±0.15 5 YHD 186478 −2.56 16 −1.32±0.05 −1.53±0.05 −1.09±0.08 −1.35±0.13 −0.72±0.16 < −0.26 −2.19±0.25 5 YHD 204543 −1.87 16 −0.63±0.05 −1.05±0.07 −0.64±0.09 −0.82±0.11 −0.21±0.16 +0.05±0.25 −1.68±0.14 5HD 206739 −1.64 3 −0.41±0.12 −0.72±0.18 · · · · · · · · · +0.20±0.26 · · · 3HD 214925 −2.08 3 −0.86±0.12 −1.09±0.20 · · · · · · · · · −0.50±0.22 · · · 3 YHD 216143 −2.32 3 −1.21±0.12 −1.24±0.15 · · · · · · · · · < −0.10 · · · 3 YHD 220838 −1.80 3 −0.76±0.12 −0.93±0.16 · · · · · · · · · +0.05±0.19 · · · 3 YHD 221170 −2.16 17 −0.73±0.06 −0.86±0.07 −0.47±0.08 −0.84±0.11 +0.02±0.13 −0.09±0.21 −1.46±0.05 17 YHD 235766 −1.93 3 −0.60±0.12 −0.86±0.14 · · · · · · · · · +0.10±0.26 · · · 3HE 1219–0312 −2.97 7 −0.75±0.07 −0.98±0.05 −0.49±0.05 < −0.89 −0.19±0.25 < +0.53 −1.19±0.21 5 YHE 1523–0901 −2.95 18 −0.63±0.06 −0.62±0.05 −0.42±0.17 −0.73±0.20 +0.24±0.05 < −0.20 −1.20±0.05 18 YM5 IV-81 −1.28 19 +0.11±0.05 −0.31±0.05 · · · −0.12±0.15 · · · +0.35±0.14 −0.58±0.15 19, 20M5 IV-82 −1.33 19 +0.11±0.05 −0.23±0.05 · · · −0.22±0.15 · · · +0.25±0.14 −0.68±0.15 19, 20M13 L598 −1.56 21 −0.34±0.07 −0.58±0.08 · · · · · · · · · +0.09±0.13 · · · 20, 21 YM13 L629 −1.63 21 −0.35±0.07 −0.61±0.08 · · · · · · · · · +0.12±0.13 · · · 20, 21M13 L70 −1.59 21 −0.23±0.07 −0.58±0.08 · · · · · · · · · +0.09±0.13 · · · 20, 21M13 L973 −1.61 21 −0.27±0.07 −0.51±0.08 · · · · · · · · · −0.01±0.13 · · · 20, 21 YM15 K341 −2.32 22 −0.73±0.08 −0.88±0.09 · · · · · · · · · · · · −1.51±0.10 22 YM15 K462 −2.25 22 −0.47±0.08 −0.61±0.09 · · · · · · · · · · · · −1.30±0.10 22 YM15 K583 −2.34 22 −1.19±0.08 −1.24±0.09 · · · · · · · · · · · · −1.70±0.10 22 YM92 VII-18 −2.29 1 −1.29±0.07 −1.45±0.07 −1.14±0.18 · · · · · · · · · −2.01±0.07 1, 2 YUMi COS82 −1.42 23 +0.52±0.16 +0.34±0.11 +0.73±0.11 · · · · · · · · · −0.25±0.15 23 Y

References. — (1) Johnson 2002; (2) Johnson & Bolte 2001; (3) Aoki & Honda 2008; (4) Cowan et al. 2002; (5) this study; (6) Sneden et al. 2003; (7) Hayek et al. 2009;(8) Christlieb et al. 2004; (9) Jonsell et al. 2006; (10) Honda et al. 2004; (11) Lai et al. 2008; (12) Hill et al. 2002; (13) Plez et al. 2004; (14) Sneden et al. 2009; (15) Cowan et al.2005; (16) Simmerer et al. 2004; (17) Ivans et al. 2006; (18) Frebel et al. 2007; (19) Yong et al. 2008a; (20) Yong et al. 2008b; (21) Yong et al. 2006; (22) Sneden et al. 2000;(23) Aoki et al. 2007a Defined by log ǫ (La/Eu) < +0.25

20 Roederer et al.

TABLE 512C/13C Ratios and Luminosities

Star 12C/13C log(L/L⊙)

BD+17 3248 6±3 2.4CS 22892–052 13±3 2.5CS 29491–069 >30 1.4HD 6268 6±2 2.5HD 74462 10±3 2.0HD 108317 15±5 1.5HD 115444 6±2 2.2HD 122563 4±1 2.6HD 122956 7±2 2.4HD 126587 10±3 2.1HD 175305 >25 1.3HD 186478 5±2 2.5HD 204543 5±2 2.5HE 1219–0312 >15 1.8

TABLE 6Pb, Th, and 12C/13C Comparisons to Previous Studies

Star Reference log ǫ (Pb) log ǫ (Th) 12C/13C

BD+17 3248 this study < +0.27 −1.26±0.10 6±3Cowan et al. (2002) < +0.3 −1.22±0.10 · · ·

CS 22892–052 this study < −0.15 −1.60±0.13 13±3Sneden et al. (2003) < −0.2 −1.59±0.10 15±2Honda et al. (2004) · · · −1.46±0.15 20

CS 29491–069 this study < +0.35 −1.46±0.25 >30Hayek et al. (2009) · · · −1.43±0.22 &90

HD 6268 this study < +0.08 −1.78±0.15 6±2Honda et al. (2004) · · · −1.97±0.10 4

HD 74462 this study +0.53±0.20 −0.94±0.13 10±3Aoki & Honda (2008) +0.35±0.26 · · · · · ·

HD 115444 this study < −0.45 −2.08±0.15 6±2Westin et al. (2000) · · · −2.27±0.11 6±1.5Johnson & Bolte (2001) · · · −2.40±0.09 6Honda et al. (2004) · · · −2.01±0.15 7

HD 186478 this study < −0.26 −2.19±0.25 5±2Johnson & Bolte (2001) · · · −2.30±0.11 6Honda et al. (2004) · · · −1.89±0.15 6

HD 204543 this study +0.05±0.25 −1.68±0.14 5±2Aoki & Honda (2008) +0.00±0.22 · · · · · ·

HE 1219–0312 this study < +0.53 −1.19±0.21 >15Hayek et al. (2009) · · · −1.29±0.14 &90

Lead and Thorium from the r-process in Metal-Poor Stars 21

TABLE 7Observed Present-Day r-process Mean Ratios

S.S. r-only “standard” stars with four “standard” four stars with anpredictionsa log ǫ (La/Eu) < +0.25 r-only starsb “actinide boost”c

Ratio (log ǫ) 〈log ǫ〉 σµ No. 〈log ǫ〉 σµ No. 〈log ǫ〉 σµ No.

La/Eu +0.179 . . . +0.18 0.03 4 +0.16 0.04 4Er/Eu +0.364 +0.40 0.03 13 +0.44 0.03 4 +0.43 0.08 4Hf/Eu −0.075 +0.10 0.04 9 +0.05 0.04 4 +0.00 0.06 1Ir/Eu +0.850 +0.88 0.04 9 +0.89 0.05 4 +0.88 0.09 2Pb/Eu · · · +0.68 0.07 7 +0.77 0.22 1 +0.17 0.15 1Th/Eu · · · −0.56 0.03 16 −0.55 0.05 4 −0.24 0.03 4

a Lodders (2003) S.S. meteoritic and Sneden et al. (2008), with updates from Gallinob BD+17 3248, CS 22892–052, HD 115444, HD 221170c CS 30306–132, CS 31078–018, CS 31082–001, HE 1219–0312

TABLE 8Comparison of Predicted and Calculated r-process Abundances for

Pb, Th, and U

206Pb 207Pb 208PbP

Pb 232Th 235U 238UP

Th, U

S.S. total (Lodders 2003) 0.601 0.665 1.903 3.169 0.0440 0.0059 0.0187 0.0686calculated r-process:

ETFSI-Q:Direct (isobaric) production 0.0209 0.0178 0.0283 0.0670 0.0184 0.0091 0.0076 0.0351Direct + indirect production 0.1439 0.1068 0.1242 0.3749 0.0415 0.0343 0.0234 0.0992

Cowan et al. (1999) 0.158 0.146 0.135 0.439 · · · · · · · · · · · ·Kratz et al. (2004) (Fe-seed) 0.163 0.151 0.138 0.452 · · · · · · · · · · · ·

r-residuals:Cowan et al. (1999) 0.240 0.254 0.158 0.652 · · · · · · · · · · · ·Beer et al. (2001) 0.178 0.171 0.133 0.482 · · · · · · · · · · · ·

Note. — The r-residuals are calculated as Nr,⊙ = N⊙ −Ns.

22

Roederer

etal.

TABLE 9Ages Derived from Individual Chronometers

“standard” stars with log ǫ (La/Eu) < +0.25 stars with actinide boostChronometer P.R. at log(P.R.) Observed No. Age Age spread Observed No. Age Age spread

pair time “zero” mean (dex) stars (Gyr) (Gyr) mean (dex) stars (Gyr) (Gyr)

Th/La 0.585 −0.233 −0.67±0.03 (σ=0.10) 16 20.4±4.1 4.7 −0.37±0.03 (σ=0.05) 4 6.4±4.3 2.3Th/Eu 0.463 −0.334 −0.56±0.03 (σ=0.08) 16 10.6±4.1 3.7 −0.24±0.03 (σ=0.06) 4 −4.4±4.3 2.8Th/Er 0.236 −0.627 −0.91±0.04 (σ=0.11) 12 13.2±4.3 5.1 −0.66±0.08 (σ=0.17) 4 1.5±5.5 7.9Th/Hf 0.648 −0.188 −0.61±0.05 (σ=0.13) 8 19.7±4.5 6.1 −0.26±0.06 1 3.4±5.7 · · ·Th/Ir 0.0677 −1.169 −1.42±0.06 (σ=0.15) 8 11.7±4.8 7.0 −1.12±0.16 (σ=0.18) 2 −2.3±8.8 8.4Th/Pb 0.111 −0.955 −1.21±0.23 (σ=0.32) 2 >9.9 11. −0.43±0.25 1 · · · · · ·

Note. — No age is derived from the Th/Pb chronometer for the stars with an actinide boost because the decay of all actinide material present could never produce the observedTh/Pb ratio in CS 31082–001.